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The OMEGA Instrument Summary Results. Bibring J-P., Langevin Y

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The OMEGA Instrument Summary Results. Bibring J-P., Langevin Y The OMEGA instrument summary results. Bibring J-P., Langevin Y., Altieri F., Arvidson R., Belluci G., Berthé M., S. Douté, Drossart P., T. Encrenaz, Forget F., Fouchet,T., Gendrin A., Gondet B., Mangold. N., Moroz V., Mustard J., Pinet P., Poulet F., Schmitt B., Sotin C., Soufflot A., Titov D., Zasova L., and the OMEGA team. 1. OMEGA instrument specifics The Observatoire pour la Minéralogie, l’Eau, les Glaces et l’Activité (OMEGA) instrument on board MEx is a visible/near infrared hyperspectral imager, built under the responsibility of IAS (Institut d’Astrophysique Spatiale), Orsay, France (Principal Investigator, PI, and Project Manager, PM) and LESIA (Laboratoire d’Etudes Spatiales et d’Instrumentation en Astrophysique), Meudon, France, in cooperation with IFSI-INAF, Roma, Italy and IKI, Moscow, Russia (Bibring et al., 2004a). OMEGA couples an imaging capability, with an IFOV of 1.2 mrad, to a spectral capability, in acquiring for each pixel the spectrum from 0.35 to 5.1 µm, in 352 contiguous spectral channels (“spectels”), with a spectral sampling of 7 nm, 13 nm and 20 nm in the spectral ranges (0.35 to 1.0 µm), (0.93 µm to 2.65 µm), and (2.51 to 5.1 µm) respectively. OMEGA includes three spectrometers, one for the visible and near infrared “VNIR” range (0.35 to 1.0 µm), and two to cover the near-infrared “NIR” range from 0.93 to 5.1 µm. Two dedicated and co-aligned telescopes are used to illuminate VNIR and NIR channels separately. VNIR operates in the pushbroom mode; it images a given line of Mars, perpendicular to the drift velocity of the spacecraft, along one direction of a bi-dimensional CCD detector array, with the spectrum of each pixel spread over the other direction of the array by a concave holographic grating. The infrared channel operates in a whiskbroom mode: one pixel of Mars is imaged at a time, and its spectrum acquired on two IR linear arrays, from 0.93 µm to 2.65 µm (“SWIR”, for “Short Wavelength IR”) and from 2.51 to 5.1 µm (“LWIR”, for “Long Wavelength IR”) respectively. A scanning mirror ahead of the NIR telescope enables contiguous pixel imaging in the cross track direction, perpendicular to the spacecraft drift velocity. For VNIR, SWIR and LWIR, the second spatial dimension is provided by the movement of the spacecraft: along the time, three dimensional (two spatial and one spectral) image-cubes are built, which constitute the OMEGA data set. The imaging performances can be summarized as follows: with an IFOV of 1.2 mrad, the spatial sampling (cross-track pixel size) is ~ 300 m when the observation is made close to periapsis, up to 4.8 km from an altitude of 4000 km. OMEGA prime goal is to perform a global coverage of Mars with a 1.5 - 4.8 km footprint, and to map selected areas amounting to ~ 5 % of the surface at high resolution (footprint < 400 m) – although the initial goal of the mission design was to acquire the global high resolution (100 m) coverage. The instrument design allows a total field of view of 128 pixels cross-track, or 8.8°, both for the VNIR pushbroom channel and the NIR whiskbroom one. However, the building of the images requires synchronizing the duration of a swath with the drift of the spacecraft in order to avoid both undersampling and oversampling. With an infrared integration time of typically 5 ms, chosen to secure a SNR > 100, the swath varies from 16 pixels when acquired at periapsis (high spacecraft orbital velocity, ~ 4 km/s), to 128 pixels when operating above 1500 km (slower spacecraft drift, < 2 km/s). Thus the high resolution mode is made of strips some 5 to 8 km wide, and hundreds of km long, while the global coverage is made of strips 300 to 500 km in width, and thousands km in length. The nominal pointing mode is nadir. In order to target given units of interest, nadir pointing with a constant cross-track offset angle must be implemented. In addition, some observations have been made in inertial mode (3 axis stabilized), in particular to acquire limb profiling, and in a “spot pointing”, or “EPF” (emission phase function) mode, to enhance the detection of atmospheric constituents. Due to severe mission resource constraints, mainly in energy and downlink capability, a strict sharing with other investigations was imposed by ESA, assigning OMEGA with an averaged data volume cap of 15%. This translated in OMEGA operating over less than one out of 4 orbits on the average, and for < 60 mn per orbit. As a consequence, a large number of specific targets were missed, and very little multiple coverage of the same area to monitor seasonal changes was performed, along the first 3 (terrestrial) years of operations. 2. Identification of atmospheric constituents As it analyzes the solar radiation diffused by Mars, OMEGA spectral images always include signatures of the atmospheric constituents, both as gases and aerosols. Their retrieval is mandatory to identify surface features, and gives access to unique atmospheric properties, with the highest spatial sampling on board Mars Express (300 m from periapsis altitude). However, given the limited spectral sampling (7 to 20 nm), only the major (CO2) and a few minor species (CO, H2O and O2) are unambiguously detected. Examples are given herebelow. 2.1. CO2 An accurate radiative transfer model devoted to the retrieval of the three components (surface, gas, aerosols) is by essence complex and iterative. At first order, it is possible to independently calculate the atmospheric absorption. A dedicated model has been built for this purpose (Melchiorri et al., 2006), which computes a multiplicative component in the synthetic spectrum representative of the main atmospheric contribution (CO2, H2O and CO) in the 1-2.7 µm range (corresponding to the SWIR-C channel). This model consists in a line-by-line spectroscopic calculation of the atmospheric absorption; a spectral atmospheric database is built for each OMEGA session, with a one-to-one correspondence. A priori knowledge of the atmospheric pressure is taken from the GCM predicted pressure (Forget et al., 1999), as well as humidity factor and CO profile. Several spectra are calculated around the GCM-predicted position in parameter space, and a least-square fit adjustment is made to retrieve the observed atmospheric pressure. Each calculation is performed at the exact illumination and elevation conditions corresponding to the OMEGA observations. The pressure measured in this way is not the actual Martian pressure, as at this stage no correction for dust opacity is performed: the “effective pressure” measured in the model is nevertheless correlated to the actual pressure, in a way that can be constrained by a more sophisticated model, including the scattering which is developed as “second order” modelling (Forget et al., 2007). Since pressure variation on Mars are primarily correlated to the altitude of the surface, CO2 being the dominant atmospheric compound, a first output of the model is an altimetry map along the OMEGA observations. This altimetry measurement can be checked against the MOLA measurement: effective pressure measurements are found to be accurate to 0.044 mbar (standard deviation), corresponding to an altimetry accuracy of 100 m (1σ). The instrumental noise itself is estimated to about 200 m. Therefore OMEGA atmospheric observation gives an altimetry measurement spatially better resolved than the a priori MOLA measurement, to an accuracy of a few hundred meters. This first step allows us to study pressure fluctuations beyond the altimetry variations: this second order study gives access to the search for global or local wave variations, at the 0.1 mbar pressure range, corresponding to the Martian meteorology. Similar searches have been tentatively made from Phobos/ISM observations (Gendrin et al., 2003). 2.2. H2O and CO variations A second step in the study of atmospheric features from OMEGA observations is the H2O and CO retrieval. Despite faint absorption at OMEGA resolution, the sensitivity of the instrument gives access to a good accuracy of the column density of these minor constituents. Study of spatial variations of the abundance ratio of H2O and CO can therefore be made, after the correction for the altitude of the surface is performed. The 2.6 µm band of H2O is used, and water vapour maps have been obtained, for Ls=94°-112°, corresponding to the sublimation of the northern -4 hemisphere polar cap. A mixing ratio of H2O in the range 2-3 10 is obtained at 40°N latitude, corresponding to 25 pr-μm, and 5-10 10-4 at 60-80°N, corresponding to 40- 60 pr-µm. These results are consistent with previous MAWD/Viking and TES/MGS below 60°N (Encrenaz et al., 2005). CO abundance is measured with only a low accuracy, due to the low intensity of the (2-0) band at 2.3 µm. Nevertheless, spatial variations have been positively observed over Hellas at Ls=130°-150°, by a factor of 2 compared to Ls=330°-350° (Encrenaz et al., 2006). This result is consistent with GCM predictions (Forget et al., 2006), which indicate an enrichment of non-condensable species over Hellas during southern winter, due to local topography affecting the global circulation. Similar variations of argon have been observed on Hellas at similar Ls (Sprague et al., 2004). 2.3. CO2 fluorescence Strong non-LTE emission of the Martian atmosphere is observed above the limb in the OMEGA observations in 3-axis stabilized mode. These emissions are identified to CO2 at 4.3 µm, with a maximum at ~ 90 km altitude and CO at 4.7 µm with a peak emission at ~ 50 km.
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